Radiation from the photosphere produces a continuous spectrum without emission or absorption lines. Above the photosphere is a layer of gas about 500 km thick, within which the effective temperature declines steadily from about 6000 to 4000K. Consequently, solar radiation has to traverse a layer of relatively cooler gas, and this absorbs radiation at wavelengths characteristic of the atoms and ions in the solar atmosphere. Within this cooler zone the absorption-line spectrum of the Sun, first described in some detail by Fraunhofer in 1814, originates. The absorption lines indicate that the mean temperature just above the photosphere is about 4500K. In older textbooks the cool zone is sometimes termed the REVERSING LAYER, and it is described as if there were a sharp transition between the hot radiating region and the overlying absorbing layer. In practice the layers merge, with continuum radiation and absorption taking place simultaneously; radiation dominates at the lower levels and absorption dominates higher up.
From the absorption-line spectrum, solar physicists have deduced the composition of the Sun’s atmosphere, and, by implication, the composition of the whole region outside the core. The elements that produce the strongest lines, such as iron, which is responsible for several hundred, are identified simply by matching features in the solar spectrum against laboratory measurements for known atoms and ions. For the fainter lines, or the less abundant elements, or the regions of the spectrum which are crowded with lines impressed by the Earth’s atmosphere (telluric absorption lines), a partly theoretical approach is necessary in order to identify the lines. Table 8.2 lists the elements which have been tracked down in the Sun by means of spectroscopy. Strong lines are produced by iron, sodium, magnesium, aluminium, calcium, titanium, chromium, and nickel.

The relative abundances of the elements in the Sun’s atmosphere are deduced from the strength of the lines which they produce, with due account being taken of the effect on the spectrum of the solar temperature and pressure. These latter exert a powerful influence: ionized calcium produces intense absorption in the visible region, whereas hydrogen gives weak lines even though it is nearly a million times more abundant. This phenomenon misled the early astrophysicists into believing that the Sun mainly consisted of heavy elements. More than two-thirds of the natural chemical elements are definitely present in the Sun. During the 1868 solar eclipse, spectroscopists observed a brilliant yellow emission line in the spectrum of the Sun’s chromosphere. Because the line could not be matched to any element known at that time, scientists attributed it to a substance called helium (from the Greek noun ‘helios’ which means. Sun). In 1895 helium was isolated in the laboratory for the first time.

Scattering of light in our atmosphere makes it impossible for us to observe the Sun’s outer atmosphere easily. Light scattered from the main solar disc is blinding in comparison to the faint glow of the solar atmosphere. During a total eclipse of the Sun, how¬ever, the Moon blocks the direct light from the disc and the Sun’s outer atmosphere is visible as a luminous halo extending several solar diameters into space. For a few seconds before and after totality of the eclipse a very thin crescent of pinkish light flashes into view. This light comes from the solar CHROMOSPHERE, a region extending from the photosphere to a height of about 2 000 km; this includes the 500 km of the reversing layer.

When the solar spectrum is observed during an eclipse, the absorption lines remain so long as part of the photosphere is visible. Once the photosphere is completely blocked, however, the absorption-line spectrum is replaced by a bright emission-line spectrum. This is the chromospheric or FLASH SPECTRUM, which is shown in figure 8.7b. It has over 3500 identifiable lines, and many of the bright ones are identical to lines in the Fraunhofer spectrum.

The temperature at the bottom of the chromosphere ia 4500k.At an altitude of 1500km or so it starts to rise to perhaps 2000 K and finally reaches 106 K where it merges with the outer¬most region, known as the corona. The height of the chromosphere actually varies and it can reach 30 (KM) km. In reality the interface between the chromosphere and the corona is not sharply definable, and therefore authorities differ as to its precise extent and temperature. The pinkish colour is mainly due to emission in the red Balmer line of hydrogen at 656.3nm. The flash spectrum also has lines of neutral and ionized helium, as well as some ionized metals that are not present in the photospheric spectrum. This shows that parts of the chromosphere are hotter than the photosphere.

The intense emission from the first Balmer line offers a method of conserving the chromosphere. The 656.3-nm line is one of the strongest absorption features in the photospheric spectrum. Therefore photographs of the Sun taken through filters which only admit light from about 656.25 to 656.35 nm record next to nothing from the photosphere (where the hydrogen is absorbing light), but a great deal of light from the chromosphere (where there is intense emission). Before the manufacture of highly-selective filters a special instrument known as a SPECTROHELIOGRAPH, invented in 1890, was used to secure photographs of the Sun in a narrow wave¬length range. Astronomers now generally use filtergrams taken in ‘Ha (656.3-nm wavelength) and the line of calcium (393.4-nm) to explore the structure of the chromosphere and its activity.

When the solar disc is photographed in either of the above lines, a network of cells, approximately the size of the supergranules, is generally visible. There may also be patches of brightness, due to high-temperature regions in the chromosphere which are produced by strong magnetic fields. High-resolution photographs also reveal many fine, dark, short lines, resembling scattered blades of grass; these are called FIBRILS. Although the fibrils are absorption features when viewed against the bright disc, they appear as luminous gas in photographs of the limb of the Sun. The fibrils look like jets or flames, and they are then termed SPICULES. They arise in the lower part of the chromosphere and may reach altitudes of 10000km. Each spicule lasts 2-10 minutes.

The chromosphere is not a uniform atmospheric layer. Essentially it is a turbulent froth churned up by the photosphere. Pressure waves generated in the convection zone are thought to be the main cause of heating in the chromosphere, which gets hotter with increasing distance from the Sun (figure 8.11). The density of matter decreases with altitude, and this causes the pressure waves to accelerate. In so doing they bring about more energetic collisions of particles which are driven along by the waves. The directed motion of the waves is thus converted to the kinetic energy of random motion of the particles, that is, to heat energy. Consequently the chromospheric temperature rises from 4300 K at its base to near-coronal values of 106 K in only 3000 km or so.

Above the chromosphere we come to the solar CORONA, a region of lo\v density and high temperature that merges imperceptibly with interplanetary space. It is difficult to observe, since its intensity in the visible region of the spectrum is only equivalent to that of the full Moon. The brilliance of the photosphere and the scattering of sunlight in our own atmosphere render it undetectable. unless the photospheric light is blocked, as happens at a total solar eclipse or in an instrument known as a CORONOGRAPH. Although Kepler observed the corona, serious interest in it did not develop until the nineteenth century. Observations of the corona during eclipses are not easy. On average, suitable eclipses occur less than annually; they are frequently over oceans or remote parts of the land; the weather cannot always be relied upon; and only a few minutes are available for obtaining data. Despite these difficulties, great efforts are made to send expeditions to. eclipse sites since certain observations can only be carried out at totality.

In the coronograph, the light from the disc is blocked off artificially by positioning an opaque disc, with diameter exactly equal to that of the Sun’s image, at the focus of the objective lens. To minimize the effects of scattered light it is essential to eliminate dust and blemishes on the objective lens and to place the coronograph at a high altitude, or even to use it above the atmosphere in a rocket or satellite. Mirrors are not normally used in coronographs because minute surface blemishes in the aluminium coating scatter too much light.

Spectroscopically , the inner corona (or K-CORONA) shows a continuum crossed by emission lines; these are fainter and less numerous than in the chromosphere. The continuum is primarily visible light from the photosphere which has been scattered towards Earth by free electrons in the corona. The Fraunhofer spectrum cannot be distinguished in the spectrum of the inner corona because the large random motions of the scattering electrons introduce large random Doppler shifts that smear out the absorption features. Beyond a few solar radii scattering of the photospheric light by dust particles takes place. These have smaller motions, and the Fraunhofer lines are not completely washed out. This region is termed the F-CORONA.

A handful of intense emission lines in the coronal spectrum intrigued the early astrophysicists. Among these lines are a strong red line at 637.4-nm wavelength and a brilliant green line at 530.3-nm wavelength; the latter was first noted by Young in 1869. These and other lines could not be matched to the characteristics of any familiar element, and so scientists speculated that an un¬known element, coronium, was responsible, just as helium had accounted for once unrecognizable photospheric lines. Later work, however, filled in the remaining gaps in the periodic table of the elements, and it became apparent that no element’s atoms could account for the coronal’lines. The resolution of the paradox came in 1941 when it was realized that familiar elements, but in a highly-ionized state, were emitting the mysterious lines. The prominent green line (530.3nm) is from atoms of iron that have been stripped of 13 outer electrons, designated Fe19f, while the red one (637.4nm) is from iron that has lost nine electrons, Few. Other lines can be attributed to Fe4+ , Fe6+, Fe8+ and Fe12+ ,(which has two intense infrared lines), and to nickel and chromium ions. There are still a few unidentified or doubtfully-identified coronal emission features. Several of the OSO (Orbiting Solar Observatory) satellites carried spectrometers. With these, the ultraviolet spectrum of the corona was observed and many new emission lines, due to highly ionized atoms, were discovered.

The high state of ionization in the corona indicates that it is at a much higher kinetic temperature than either the photosphere or chromosphere .An enormous amount of energy is needed to dislodge 13 orbital electrons from a neutral iron atom and so form Fe13+ ions. Collisions between atoms and ions only become powerful enough to strip these electrons at temperatures of about 106K; occasionally a temperature of 2 X 106K is reached in the outer corona, and individual hotspots may temporarily reach 4 x 106K. Because the coronal gas is rarefied (the density is 1011 particles m”3, compared to 1025 particles nr3 in Earth’s atmosphere) the total amount of energy stored in it is not great. Even so the in¬tense heating of this tenuous gas has long puzzled solar physicists. The continual dumping of energy into the corona by shock waves and turbulence arising at the photosphere could be the most significant source of coronal heating. Solar flares, and other forms of activity discussed below, may be contributory agents.

Much of the radiation from the corona falls beyond the visible region in the extreme ultraviolet and X-ray domains of the spec¬trum. This is because the corona is very hot and highly ionized. X-ray images of the Sun show areas of locally higher particle density and higher temperature in the inner corona. These hotspots are generally located high above active regions on the surface of the Sun. The detailed structure evident in such photographs proves that the corona is non-homogeneous.

The corona is in dynamic rather than static equilibrium, as it expands under its own pressure gradient, against the Sun’s gravitational field, into the near-vacuum of interplanetary space. The flow of material out of the corona, called the solar wind, takes about five days to reach Earth. It has a considerable effect on the magnetic fields of the planets since the ionized matter of which it is composed cannot cross magnetic field lines. In the neighbourhood of planetary magnetic fields, the wind flows round a protective magnetic shell.

Comments are closed.